Ly alpha emitting galaxies as early stages in galaxy formation

Lennox L. Cowie,4 Amy J. Barger,567
Esther M. Hu8

Abstract

We present optical spectroscopy of two samples
of GALEX grism selected Lyα emitters (LAEs):
one at z=0.195−0.44 and the other at z=0.65−1.25.
We have also observed a comparison sample of galaxies
in the same redshift intervals with the same UV magnitude distributions
but with no detected Lyα.
We use the optical
spectroscopy to eliminate active galactic nuclei (AGNs) and to
obtain the optical emission-line properties of the
samples.
We compare the luminosities of the LAEs in the two
redshift intervals and show that there is dramatic evolution
in the maximum Lyα luminosity over z=0−1.
Focusing on the z=0.195−0.44 samples alone, we show that there
are tightly defined relations between all of the galaxy parameters
and the rest-frame equivalent width (EW) of Hα. The
higher EW(Hα) sources all have lower metallicities,
bluer colors, smaller sizes, and less extinction, consistent
with their being in the early stages of the galaxy formation
process. We find that 75±12% of the LAEs
have EW(Hα)>100 Å, and, conversely, that 31±13%
of galaxies with EW(Hα)>100 Å are LAEs.
We correct the broadband magnitudes for
the emission-line contributions and use spectral
synthesis fits to estimate
the ages of the galaxies. We find a median age of 1.1×108 yr
for the LAE sample and 1.4×109 yr for the UV-continuum
sample without detected Lyα. The median
metallicity of the LAE sample is 12+log(O/H)=8.24, or
about 0.4 dex lower than the UV-continuum sample.

Lyα emission-line searches have been widely used to find
high-redshift galaxies and, for the highest redshift galaxies, this
line is the only spectroscopic signature that can be used to confirm
the redshift of a galaxy selected on the basis of its color properties.
However, Lyα is a difficult line to interpret.
Because the line is resonantly scattered
by neutral hydrogen, determining its escape path and
hence its dust destruction is an extremely complex problem,
both theoretically (e.g., Neufeld 1991; Finkelstein et al. 2007)
and observationally (e.g., Kunth et al. 2003; Schaerer & Verhamme 2008;
Östlin et al. 2009).
Thus, while we have empirical measurements that a significant fraction of
UV-continuum selected samples have Lyα lines
with rest-frame equivalent widths above 20 Å over a wide range of redshifts from z=0.3 to z=6.5
(Shapley et al. 2003; Cowie et al. 2010, 2011; Stark et al. 2010),
our understanding of what determines this fraction is still weak.
In particular, we would like to know whether the presence of Lyα
emission is related to other properties of the galaxy, such as its
metallicity, extinction, morphology, or kinematics,
and how the Lyα line escapes.

Due to the difficulty with observing in the UV, we currently have much more
information on the z∼2−3 Lyα emitters (LAEs) and how their
properties relate to those of other UV selected galaxies at these redshifts
(e.g., Shapley et al. 2003; Reddy et al. 2010; Kornei et al 2010)
than we do on the local samples. However, there has been considerable controversy
in the interpretation of these high-redshift observations. The simplest
interpretation is that the LAEs are younger, lower mass, and metal poor,
representing early stages in galaxy evolution (e.g., Hu et al. 1998;
Nilsson et al. 2007; Gawiser et al. 2007).
However, other authors have argued that LAEs arise in
relatively massive galaxies (e.g., Lai et al. 2008; Finkelstein et al. 2009c)
with ages of around a Gyr, that young galaxies have weaker Lyα than old
galaxies (Shapley et al. 2001), and, more recently, that LAEs are older, less
dusty, and in a later stage of galaxy evolution than sources with weaker
Lyα emission (Kornei et al. 2010).

The Lyα signature can be produced by a range of galaxy types and
masses, including even the most ultraluminous infrared galaxies (ULIRGs)
(e.g., Chapman et al. 2005; Nilsson & Møller 2009), so some level
of heterogeneity must be expected. However, all of the results that
argue for the LAEs being predominantly old are based on spectral synthesis
fitting, and most of the old ages inferred are almost certainly mis-estimates
arising from the presence of very strong optical emission lines in the
LAEs (e.g., Schaerer & deBarros 2009).

Until recently the only low-redshift Lyα emitting sources that
could be studied in detail were the local blue compact galaxies.
However, these generally have much lower Lyα luminosities than
the high-redshift LAEs, and, while some of the blue compact galaxies have
been studied in exquisite detail (e.g., Östlin et al. 2009) on an
individual basis, it has not been easy to form large, uniformly selected
samples that can be statistically analyzed. Thus, the recent determination
that substantial z∼0.2−0.4 samples of LAEs can be found
(Deharveng et al. 2008) with the Galaxy Evolution Explorer (GALEX)
(Martin et al. 2005) grism spectrographs has
enabled a new approach to the subject (Atek et al. 2009a;
Finkelstein et al. 2009a, 2009b; Scarlata et al. 2009; Cowie et al. 2010).

The low-redshift LAE samples have many
advantages. The galaxies are bright and can be easily
studied at other wavelengths, but perhaps even more
importantly, they can be integrated into comprehensive
studies of galaxies at the same redshifts to understand the
selection biases intrinsic to the samples.
Early papers on GALEX LAEs worked with relatively small samples,
but the general conclusions are that low-redshift LAEs are somewhat
heterogeneous yet more weighted to low metallicities and extinctions
and more likely to be small, compact galaxies when compared to
UV-continuum selected galaxies without detected Lyα with the
same luminosities in the same redshift interval.

Here we study larger and more optically spectroscopically complete
samples of LAEs at z=0.195−0.44 and z=0.65−1.25, together with
comparison samples of UV-continuum selected galaxies without detected
Lyα in the same redshift intervals.
In Section 2 we present our optical spectroscopy of
all the samples obtained with the DEep Imaging Multi-Object Spectrograph
(DEIMOS; Faber et al. 2003) on the Keck II 10 m telescope.
In Section 3 we use the optical data to remove the small
number of active galactic nuclei (AGNs) that were not previously
identified from the UV spectra and then provide our final sample of
candidate LAE galaxies in the two redshift intervals. We also use
the spectra to measure the optical line fluxes and equivalent
widths and to determine the metallicities of the galaxies.
In Section 4 we determine how the Lyα
luminosity evolves with redshift.
In Section 5 we consider the overall properties
of the LAEs and the UV-continuum selected galaxies without detected
Lyα.
Since the sample of z∼1 LAEs is small (only eight objects)
we only consider the z=0.195−0.44 sample in this section.
We use the spectra to determine
the emission-line contributions to the broadband
fluxes and to show that these corrections must be included
if spectral synthesis fitting is to give accurate
ages for the youngest (<109 yr) galaxies.
In Section 6 we interpret the results in terms of a constant
star formation rate (SFR) model, comparing SFRs derived from the Hα,
UV-continuum, and 20 cm observations and discussing the limits of
validity for the UV-continuum determined SFRs. We also interpret the
metallicity evolution.
We summarize our results in Section 7.
We use a standard H0 = 70 km s−1 Mpc−1,
ΩM = 0.3, ΩΛ = 0.7
cosmology.

For the present study we use the LAE samples with rest-frame
EW(Lyα)>15 Å from
Cowie et al. (2010; their Tables 15 and 16), chosen from the nine
blank high galactic latitude fields with the deepest GALEX
grism spectroscopic observations (Morrisey et al. 2007).
The area covered is just over 8 deg2. The samples consist of
sources whose grism UV spectra have a detectable single emission
line, which we assume to be Lyα. In order to eliminate
sources that are clearly AGNs based on their UV spectra, we do
not include any sources that have high-ionization lines, and we
require the FWHM of the Lyα lines to be less than 15 Å in the GALEX FUV spectra (z=0.195−0.44) or less than
30 Å in the GALEX NUV spectra (z=0.65−1.25).
In the redshift interval z=0.65−1.25 we also eliminate sources
that do not show the break between the FUV and NUV bands. Such a
break would be expected to be produced by the Lyman continuum
edge. (See Section 4 of Cowie et al. 2010 for details on
constructing the samples.)
For the CDFS 00 field we have added additional sources
from a deeper grism spectroscopic exposure than was used in
Cowie et al. (2010). The full Lyα
selected sample in the CDFS 00 is given in the Appendix
(i.e., this is an update of Table 4 in Cowie et al. 2010, but
it also includes the optical redshifts, where available; see
Section 2.2).

We summarize the final candidate LAE galaxy sample in the redshift
interval z=0.195−0.44 in Table 1, sorted
by the rest-frame EW(Lyα).
Through most of the paper we restrict
our analysis to sources with rest-frame EW(Lyα)≥20 Å, so
we list these sources in the main body of the table.
However, for completeness, we also give the results for
the lower EW(Lyα) sources in a supplement to the table.
For each source in Table 1 we give the GALEX
name, the J2000 right ascension and declination, the
redshift from the GALEX UV spectrum,
the logarithm of the bolometric continuum luminosity, Lνν, above
1216 Å, the logarithm of the Lyα luminosity, the
rest-frame EW(Lyα), the optical ground-based redshift,
the rest-frame EW(Hα), the [NII]λ6584/Hα,
[OIII]λ5007/Hβ, and [OIII]λ4363/Hγ ratios,
and the logarithm of the Hα and Hβ fluxes for the sources
with optical continuum magnitudes in the Sloan Digital Sky Survey
(SDSS) (see Section 2.1).

We summarize the final candidate LAE galaxy sample in the redshift
interval z=0.65−1.25 in Table 2, also sorted by the rest-frame
EW(Lyα). There are 8 sources in this sample. We list
the properties of these sources using the same format as in Table 1,
except here we give the rest-frame EW(Hβ) in place of the rest-frame
EW(Hα), and we do not include the column for the Hα and
Hβ fluxes. We also do not split the sample into a main body of
the table and a supplement to the table, since there are so few
sources. Six of the eight sources satisfy the EW(Lyα)>20 Å criterion.

Although the z∼1 sample is small, it is important in that it
demonstrates the presence of these LAEs, and it allows us to
make a first estimate of the
luminosity function at these redshifts. However, most of our
analysis in this paper will focus on the z∼0.3 sample.
For each source with z<0.3 we computed the UV spectral index
from the GALEX spectra using the Calzetti et al. (1994)
prescription, which fits a power law to the spectrum
in a set of wavelength windows chosen to
minimize the effects of absorption and emission lines.
Since we have only short wavelength data in the
rest frame, we corrected our measured index with the offset of −0.16
given by Meurer et al. (1999) to provide our final index. We did
not compute the UV spectral index for the sources with z>0.3, since
at those redshifts many of the shortest wavelength windows in
Calzetti et al. (1994) lie in the noisy regions between the
GALEX FUV and NUV spectra.

2.1 Imaging Data

We compiled ancillary data from various archival sources. Just over
82% of the area is covered by deep U-band images obtained with
the MegaPrime camera on the 3.6 m Canada France Hawaii Telescope (CFHT).
We took the reduced images from the CADC
pipeline reduction, which gives the 5σ AB magnitude limits.
These range from 26.5 to 27.5. We used these images to measure the
U-band magnitudes and the FWHM sizes of the galaxies at this
wavelength. For the fields covered by the SDSS we
also compiled the u′,g′,r′,i′, and z′
SDSS model C magnitudes (2σ AB magnitude limits of 22.0, 22.2,
22.2, 21.3, and 20.5, respectively)
from the DR6 release
(Adelman-McCarthy et al. 2008). These cover
roughly half of the observed targets. For the fields
with publicly available 20 cm images
(COSMOS 00: Schinnerer et al. 2007, 1σ of 10.5μJy beam−1;
SIRTFFL 01: Condon et al. 2003, 1σ of 23μJy beam−1;
HDFN 00: Morrison et al. 2010, 1σ ranges from 3.9μJy beam−1
near the center to 8μJy beam−1 at 15′),
we measured the 20 cm fluxes from the images following the usual
methods (e.g., Morrison et al. 2010).
Finally, for the CDFS 00 targets covered by the
Galaxy Evolution from Morphologies and SEDs (GEMS)
survey (Rix et al. 2004), we compiled thumbnails of the
galaxies in the F606W (5σ AB magnitude limit of 28.25) and
F850LP (5σ AB magnitude limit of 27.10) filters from
the Hubble Space Telescope (HST) observations with the
Advanced Camera for Surveys (ACS).

2.2 Optical Spectroscopy

We obtained optical spectroscopy for a large fraction of the two
GALEX LAE samples using the DEIMOS spectrograph on Keck II
with the ZD600 ℓ/mm grating
in observing runs throughout 2009 and 2010.
The ∼0.5−1.0 hr spectra have a resolution of 4.5 Å and a wavelength range of approximately 5000 Å centered at 6800 Å.
More details of the observing and reduction procedures can be found in
Cowie et al. (1996, 2010).

We observed 77 of the 91 z=0.195−0.44 LAEs in Table 1
(this includes the supplement) and all 8 z=0.65−1.25 LAEs in
Table 2.
We give the optical redshifts in Tables 1
and 2, where they can be compared with the UV redshifts from
the GALEX data. For the 77 z=0.195−0.44 LAEs with optical
redshifts, we confirmed the UV redshifts for 70.
In Table 1 we show the optical redshifts in parentheses
for the sources where we did not confirm the UV redshifts.
We eliminated these sources from further consideration.
We confirmed the UV redshifts for all 8 z=0.65−1.25 LAEs
with the optical redshifts.
We show the UV and optical spectra for these sources
in Figures 2.2 and 2.2.

In order to understand how the GALEX LAEs
relate to other UV-continuum selected sources with the same
UV-continuum magnitude distributions, we also observed at optical
wavelengths sources randomly chosen from the GALEX UV
spectroscopic sample that did not have UV spectral line identifications.
The latter is not a complete sample
of all such sources in the field, but rather a randomly chosen
subsample with an NUV magnitude selection function that is similar
to that for the sources with strong Lyα emission. Thus,
it may be directly compared to the LAE sample.

We have obtained optical spectra for 450 of
the 5642 objects in the total UV-continuum selected
sample without detected Lyα. 114 of these
sources lie in the redshift interval z=0.195−0.44,
and 5 lie in the redshift interval z=0.65−1.25.
We will hereafter refer to these sources as our UV-continuum sample.
These sources may be viewed as analogs of high-redshift Lyman
break galaxies (LBGs) without strong Lyα emission and
can be combined with the LAE sample to understand how
Lyα galaxies are related to the more general
UV-continuum selected population.
In order to compare the numbers of UV-continuum sources with a given
property with the LAE sample we need to correct for the fraction
of UV-continuum sources that were observed.
There is a weak dependence of the observed fraction on the
NUV magnitude, ranging from 6% observed at NUV=20 to 11%
observed at NUV=22. We use the
inverse of this fraction as a function of the NUV magnitude
to provide the weighting to the UV-continuum sample.

The new optical spectroscopic data presented in this paper
more than doubles the optical spectroscopic data used in
Cowie et al. (2010).
There we observed with DEIMOS (including four literature redshifts)
34 of the 62 z=0.195−0.44 LAEs given in their Table 15 (this also
includes EW(Lyα)<20 Å sources) and 1 of the 4
z=0.65−1.25 LAEs given in their Table 16. We
also observed with DEIMOS 124 UV-continuum selected sources
without detected Lyα, of which 46 were in the redshift
interval z=0.195−0.44 and none were in the redshift interval
z=0.65−1.25.

Figure 1: GALEX spectra of the first four z=0.65−1.25
LAEs in Table 2.
For each source we give the redshift measured from the UV spectrum.
The Lyα line is marked, as are the positions where CIV and NV
would fall. The blue vertical dashed line shows the position of
the Lyman continuum edge at the galaxy redshift.
This is generally slightly redward of the more accurate optical
redshift (see Cowie et al. 2010).

Figure 2: DEIMOS optical spectra of the first four z=0.65−1.25 LAEs
in Table 2. For each source we give the redshift measured from the
optical spectrum. The positions of the [OII]λ3727 and [OIII]λ5007
lines are marked. Where continuum magnitudes are available
we have normalized the spectrum to match these. None was
available for the source in (b), so its normalization should be
considered to be much more uncertain.

Figure 1 (cont): GALEX spectra of the last four z=0.65−1.25 LAEs
in Table 2.
For each source we give the redshift measured from the UV spectrum.
The Lyα line is marked, as are the positions where CIV and NV
would fall. The blue vertical dashed line shows the position of the
Lyman continuum edge at the galaxy redshift.
This is generally slightly redward of the more accurate optical
redshift (see Cowie et al. 2010).

Figure 2 (cont): DEIMOS optical spectra of the last four z=0.65−1.25 LAEs
in Table 2.
For each source we give the redshift measured from the optical spectrum.
The positions of the [OII]λ3727 and [OIII]λ5007
lines are marked. Where continuum magnitudes are available
we have normalized the spectrum to match these. None
was available for the source in (f), so its normalization should be
considered to be much more uncertain.

2.3 Line Fluxes

Most of the spectra were obtained at a near-parallactic angle
to minimize atmospheric refraction effects. However,
in some cases the mask configuration did not allow for this.
In addition, a small number of the spectra were obtained during periods of
varying extinction.
All of the spectra were relatively calibrated using
the measured response from observed
calibration stars. However, relative slit losses always pose a
problem, and special care must be taken for the flux calibration.

For each spectrum we fitted all of the emission lines using the
IDL MPFIT program of Markwardt (2009). We used simultaneous Gaussian
fits to neighboring lines together with a linear fit to the continuum
baseline. For weaker lines we held the full width constant, using
the value measured in the stronger lines, and set the central wavelength
to the measured redshift. We also measured the noise as a function of
wavelength by fitting to random positions in neighboring portions of the
spectrum and computing the dispersion in the results.

For the AGN classifications and the metallicity measurements we used
only the ratios of the emission lines that are close in wavelength.
These line flux ratios can be robustly measured from the spectra without
even the relative flux calibration, since for these neighboring lines the
DEIMOS response is essentially constant.
For the primary lines used in our analysis we give the
line ratios and their 1σ errors in Tables 1 and 2.

Figure 5: Comparison of the f(Lyα)/f(Hα) ratios
(red squares) and the f(Hβ)/f(Hα) ratios (blue diamonds)
measured from the present data with the ratios measured by
Scarlata et al. (2009). The scale for the f(Lyα)/f(Hα)
ratios is shown on the bottom and left axes and that for the
f(Hβ)/f(Hα) ratios on the right and top
axes. For the f(Lyα)/f(Hα) ratios
the vertical error bars only include the uncertainties in
f(Hα), since the uncertainties in f(Lyα) are
not given in Scarlata et al.’s table.

In order to measure the absolute Hα and Hβ line fluxes, we
used the measured EW(Hα) and EW(Hβ) in combination with
the measured continuum fluxes at the appropriate wavelength.
We first determined which SDSS filter the emission line lay in and
then integrated through the transmission of this filter
to obtain the averaged spectral flux. We then renormalized
the spectrum to match the corresponding SDSS model C
magnitude from the DR6 release
(Adelman-McCarthy et al. 2008). We then determined
the continuum flux and multiplied this by the
observed-frame EW to obtain the corresponding line flux. This
is an approximation, since it assumes that the measured
EW is representative of the value averaged over the total
light of the galaxy, including regions outside the slit.
However, for the photometric cases we derive very similar
values directly from the calibrated spectra, including
calibrated spectra that we obtained using larger
(2′′) slit widths for a small subset of the galaxies.
This suggests the procedure is relatively robust.

However, the relative apertures used in measuring the GALEX
Lyα fluxes and the optical emission line fluxes may introduce
systematic variations reflecting the relative geometries
of the emission. Thus, the f(Lyα)/f(Hα) ratios
are more reliably obtained in large (5′′)
aperture measurements, such as those used in Atek et al. (2009a).
In the present paper we are primarily concerned with determining
the optical continuum and emission-line properties that result
in a galaxy showing a strong Lyα emission line and not
so much with the relative fluxes of the optical and Lyα
emission lines, making this less of a concern. In the discussion
we will note where aperture effects might play a role.

In order to check the calibrations and aperture effects,
we tested our measured fluxes against the independent measurements
by Scarlata et al. (2009) for the subset of galaxies that overlap
with our sample. The Scarlata et al. data were obtained with a
larger 1\farcs5 slit. For eight cases
they compared their results with those made in a 5′′ slit, finding
only small differences between the two measurements.
Scarlata et al.’s procedure for measuring the fluxes relied on the
relative accuracy of the spectrophotometry, using the i′ magnitude
from the SDSS data to set the absolute flux normalization.
This differs from our procedure, where we normalized
each region of the spectrum separately to the
SDSS magnitudes. We compare the two results in
Figure 2.3, where we show
both the f(Lyα)/f(Hα) (red squares)
and the f(Hβ)/f(Hα) ratios (blue diamonds).
The overall agreement is very good, but in a small number of
cases the disagreement is much larger than the statistical
error in the f(Lyα)/f(Hα) ratio. In
the most extreme case the ratios differ by a multiplicative
factor of just over two. We take this to be an estimate of the
potential systematic errors.

We might expect that sources classified as AGNs based on
the presence of high-excitation lines or very wide lines in the UV
spectra are truly AGNs. However, the converse
is not true, which means our candidate LAE galaxies (Tables 1
and 2) may still have some degree of AGN contamination.
In some cases the high-excitation lines fall at problematic
wavelengths in the UV spectra, and in other cases these lines may be
intrinsically weak. Optical spectroscopic data support these points:
follow-ups of GALEX candidate LAE galaxy samples have shown
that a significant fraction (∼20%) have substantial
AGN contributions (Finkelstein et al. 2009b; Scarlata et al. 2009; Cowie et al. 2010), though there is wide variation
in the estimates of the precise degree of contamination
reflecting the relatively small sample sizes and
the differences in the classification procedures.
(Finkelstein et al. 2009b find a higher AGN fraction in their sample
than other groups, so we make a detailed comparison with their results
in this section.)

We first searched our samples for optical spectra that showed
broad emission lines. None of the spectra are quasars or Seyfert 1s,
a consequence of this class of source being easily picked out
with the UV spectra (Barger & Cowie 2010). However, three of the sources
in the z=0.195−0.44 sample are intermediate-type Seyferts.
In the optical redshift column (col. 8) of Table 1 we have labeled
these sources as Sy1.5 or Sy1.8. We eliminate these sources from further
consideration.

We next constructed the BPT diagram (Baldwin et al. 1981)
for the z=0.195−0.44 sample. This diagram uses the [NII]λ6584/Hα
and the [OIII]λ5007/Hβ ratios to separate AGNs from star-forming
galaxies. Hα is off the spectra in the z=0.65−1.25 sample, and we
cannot run this test there. In addition, only sources with strong optical
emission lines can be classified in this way. We restricted to galaxies
that had either [NII]λ6584 or Hα detected with a
signal-to-noise above 5 and either [OIII]λ5007 or Hβ
detected with a signal-to-noise above 5.

The BPT diagrams for the (a) LAE and (b) UV-continuum
samples having strong optical emission lines defined in this way
are shown in Figures 3(a) and 3(b), respectively.
The dotted curve traces the Kewley et al. (2001) division between
galaxies whose extreme UV ionizing radiation field is dominated by
an AGN (>50%) and those dominated by star formation. It is a
theoretical curve based on photoionization models for giant HII regions
and a range of stellar population synthesis codes. The dashed
curve traces the Kauffmann et al. (2003) division between AGNs and
star formers based on an empirical separation of SDSS galaxies.
As mentioned above, we have eliminated the intermediate-type Seyferts
in Table 1 from Figure 3(a), but the BPT diagram
suggests that a further four sources are also AGNs.
We mark these with open black squares in Figure 3(a) and
label them as “B” in the optical redshift column (col. 8) of Table 1.
Another four sources have weak [OIII] and Hβ
and do not appear on the diagram, but they
have strong log([NII]λ6584/Hα)>−0.25 ratios,
which suggests that they are also AGNs.
We also label these sources as “B” in column 8 of Table 1.
Most of the remaining sources in Figure 3(a) clearly lie along
the star-forming galaxy track, as do all the UV-continuum
sources in Figure 3(b), though there are a small number
of sources that lie in positions where they may
have mixed star formation and AGN contributions.
None of the results in the paper are changed if we
exclude these objects.

Figure 6: (a) BPT diagram for the candidate LAE galaxy sample
(Table 1; red solid squares). Sources from both the main table
and the supplement (the latter contains sources with rest-frame EW(Lyα)
between 15 and 20 Å) are included. Only the 58 sources with strong optical
emission lines (as defined in the text) are shown.
The intermediate-type Seyferts in Table 1 are excluded.
The black open squares enclose the sources that we classify as
AGNs on the basis of this diagram.
(b) BPT diagram for the sources in the
UV-continuum sample (blue solid diamonds). Again, only
the 82 sources with strong optical emission lines are shown.
In both panels the green diamonds show sources where the
[OIII]λ4363 auroral line is detected above the 3σ level,
the dashed curve denotes the Kauffmann et al. (2003) empirical division
between AGNs (upper right) and star-forming galaxies (lower left), and
the dotted curve denotes the Kewley et al. (2001) theoretical division.

There remain 8 sources with optical spectra in the LAE
sample where the [OIII]λ5007
and Hβ lines are too weak to place them on the BPT diagram
and where the [NII]λ6584/Hα ratios are not high.
These galaxies are labeled as “A” in the optical redshift
column (col. 8) of Table 1 for their weak optical emission lines.
In one case the optical spectrum is poor, but in the remaining cases
the sources either have more normal [NII]λ6584/Hα ratios or are
absorbers. It is possible that these also contain AGNs.
The percentage of AGNs in the sample is 16±4% if we include only the
classified AGNs, but that percentage would rise to 27±6% if the 8
weak emission-line sources were also AGNs.

Finkelstein et al. (2009b) optically classified 23 sources from the
Deharveng et al. (2008) sample in the Groth strip. Fourteen of
these overlap with sources in our Table 1. The remaining 9 do not,
either because the source lies outside the region of the field we used or outside the
redshift range we used, or because we found the Lyα identification to be
dubious, or because we classified the source as an AGN based on the UV spectrum.
Of the 14 sources in common, the classifications agree for 11 (3 AGNs and
8 star formers). We could not classify GALEX1420+5306, the first of the remaining
3 overlapped sources, as it has very weak emission lines.
Finkelstein et al. (2009b) found this source, which they called EGS1,
to be an AGN based on its [NII]λ6584/Hα
ratio. As discussed in Cowie et al. (2010), we believe the second of
the remaining 3 overlapped sources,
GALEX1417+5224 (EGS2), is a high-excitation, very low metallicity galaxy
rather than an AGN. Finally, we find that the third source, GALEX1421+5239 (EGS14),
lies on the star-forming galaxy track in the BPT diagram. We do not
detect high-ionization lines. This is in contrast to Finkelstein et al. (2009b),
who classified the source as an AGN on the basis of both its position on
the BPT diagram and the presence of high-ionization lines in its spectrum.

All eight of the z=0.65−1.25 LAEs in Table 2 have been optically
observed (Figures 2.2 and 2.2),
and in all cases the Lyα identification is confirmed with the
optical redshift. Seven of these have strong emission
lines, and none of them shows [NeV] emission or broad MgIIλ2800
emission that might classify the source as an AGN. The remaining source
(GALEX1420+5313; see Figure 2.2(e)) has a post-starburst
spectrum. (This galaxy is also identified in the DR3 release of the DEEP2
survey; Davis et al. 2007.) We label it as “A” in column 8 of
Table 2 for its weak optical emission lines.
It is possible that the Lyα
emission is AGN-dominated in this galaxy. We therefore
exclude this source from the z=0.65−1.25 LAE sample.

Throughout our subsequent analysis we restrict to the LAEs that are confirmed
as star formers from the optical spectra. We eliminate all sources with AGN
signatures in either the optical or UV (including those sources identified as
AGNs based on the BPT diagram), and we also eliminate the optically
unobserved galaxies, the unconfirmed galaxies, and the galaxies
with weak optical emission lines. This leaves a sample of 40 LAEs
with rest-frane EW(Lyα)≥20 Å in the z=0.195−0.44 interval and a sample of
5 LAEs with EW(Lyα)≥20 Å in the z=0.65−1.25 interval.

Figure 7:
The average UV spectra of the 40 z=0.195−0.44 LAEs (red spectrum)
and of the 101 UV-continuum sources with rest-frame
EW(Hα)>5 Å (blue spectrum). The gaps in the spectra
correspond to the wavelength region between the FUV and NUV
GALEX grisms. In addition to Lyα, we mark
the positions of UV absorption features that are only weakly
seen in these low-resolution spectra.

In Figure 3 we compare the average UV
spectrum of the 40 z=0.195−0.44 LAEs (red spectrum)
with that of the UV-continuum sample
in the same redshift interval (blue spectrum).
In order to make the most direct comparison, we
eliminate weak emission-line galaxies with rest-frame
EW(Hα)<5 Å in the UV-continuum
sample, since we have eliminated these sources
from the LAE sample. We also eliminate two objects
which are classified as BPT AGN in the UV-continuum sample.
This leaves 101 UV-continuum
galaxies. However, including the remaining
galaxies makes no difference to the results.

The NUV normalizations are almost identical since the samples are
chosen with the same distribution of NUV magnitudes. The shapes
are also nearly indistinguishable above 1400 Å. By selection
the LAE sample has strong emission, and the average spectrum for
this sample has a rest-frame EW(Lyα)=36 Å. While a
a stacking procedure is not the best way to analyze objects
with a mixture of Lyα emission and absorption, it
at least allows us to see that the UV-continuum sample has
Lyα emission and that its stacked spectrum is
flatter than that of the LAE sample below 1400 Å.
Fitting the Lyα emission, we find a rest-frame
EW(Lyα)=7 Å for the emission
feature. This is comparable to, but slightly higher than,
the median value of EW(Lyα)=4 Å found for
the z∼3 LBG population by Kornei et al. (2010). The z∼3
LBG sample includes the LAEs as a subsample; thus, if anything, our
sample, which excludes the LAEs, should be lower.

A rest-frame EW(Lyα)≥20 Å is normally used to define
the high-redshift LAE population (e.g., Hu et al. 1998), so in
this section we apply that criterion.
In Figure 4 we plot Lyα luminosity versus
redshift and compare it with the measured luminosities of high-redshift LAE samples
taken from Shimasaku et al. (2006), Ouchi et al. (2008), and Hu et al. (2010),
all of which were chosen with the same EW selection criterion.

The primary conclusion that we draw from Figure 4
is that there is a dramatic change in the properties of
the LAEs between z=1 and z=0. There are
no sources in the z=0.195−0.44 sample with luminosities
close to those of the brightest LAEs at high redshifts,
but such sources are present by z∼1.
This is not simply a volume effect: Cowie et al. (2010)
determined the LAE luminosity function in the z=0.195−0.44 redshift
interval and showed that the L⋆ for a Schechter (1976)
function fit with α=−1.36 (red dashed line) is a factor of
8 fainter than that measured by Gronwall et al. (2007) at z∼3
and more than an order of magnitude fainter than that
measured by Hu et al. (2010) at z=5.7 (blue solid line).

We have computed the z=0.65−1.25 Lyα luminosity
function for the observed sample using the 1/V technique
(Felten 1976) with the accessible volumes calculated from
the areas versus NUV magnitudes in the sample. We show
the luminosity function with the open black squares in
Figure 4. As discussed extensively in
Deharveng et al. (2008) and Cowie et al. (2010),
this Lyα luminosity function corresponds to sources
with NUV<22 and must be corrected for LAEs
at fainter continuum magnitudes. In the redshift interval
z=0.195−0.44 most sources with
Lyα luminosities in the 1042−1043 erg s−1
range lie at NUV<22 and the corrections are
small (Cowie et al. 2010). In contrast, many of the LAEs
wth Lyα luminosities
∼1043 erg s−1 at z=0.65−1.25
are fainter than NUV>22 and the correction is substantial.
Following the procedures outlined in Deharveng et al. (2008)
and Cowie et al. (2010), we have computed the correction by
assuming the z∼1 NUV counts derived from
the luminosity functions of Arnouts et al. (2005)
and the shape of the equivalent width
distribution derived in Cowie et al. (2010)
for the z=0.195−0.44 sample. (This shape is
almost identical to that seen at z∼3
by Shapley et al. 2003.) We show the corrected points
with the solid squares in Figure 4.

We compare with the Lyα luminosity
functions at z∼0.3 (blue curve),
z∼3 (red curve), z∼5.7 (green solid curve),
and z∼6.5 (green dashed curve) using the Schechter (1976)
function fits given in Cowie et al. (2010),
Gronwall et al. (2007), and Hu et al. (2010). Even allowing
for the small numbers in the present sample and the substantial
corrections, we can see that
there is a dramatic evolution in the luminosity
function between z∼0.3 and z∼0.95.
The data do not justify a Schechter (1976) function fit,
but we show the Gronwall et al. (2007) luminosity function
with L⋆ reduced by a factor of 2.5 as
the black dashed curve. This function
(ϕ⋆=1.28×10−3, α=−1.36, and
logL⋆=42.26) provides a reasonable fit
to the data, though it is by no means unique. It
would imply a reduction by a factor of 2.5 in the
Lyα luminosity density between z=3 and z=1,
which is much smaller than the dramatic factor of 50 reduction
that Cowie et al. (2010) estimate between z=3 and z=0.3.
This suggests that much of the evolution occcurs
in the z=0−1 range, as is seen in many other populations.

Figure 9: The derived Lyα luminosity function at
z=0.65−1.25 (black squares). The open squares show
the values calculated for the present sample,
which is drawn from galaxies with NUV<22.
The solid squares show the values corrected to
allow for sources with fainter continuum mangitudes
(see text). The error bars are ±1σ from the Poisson
errors corresponding to the number of sources in the bin.
The blue curve shows the Schechter (1976) function
fit to the Lyα luminosity function at z=0.194−0.44
from Cowie et al. (2010), the red curve shows the fit
at z=3 from Gronwall et al. (2007), and the green curves
show the fits at z=5.7 (solid) and z=6.5 (dashed)
from Hu et al. (2010). The black dashed
curve shows the Gronwall et al. luminosity function with
L⋆ reduced by a factor of 2.5
to fit the z=0.65−1.25 points.

5.1 Overview

We now turn to the z=0.195−0.44 samples alone and what they can
tell us about how LAE galaxies are drawn from the more general
UV-continuum selected galaxy population. In this section we will
occasionally include the LAEs with EW(Lyα)<20 Å from the
Table 1 supplement when this is appropriate. We parameterize
the data with the rest-frame EW(Hα). Using the EW(Hα)
allows us to intercompare easily the properties of the
Lyα and UV-continuum samples. In addition, the EW(Hα)
is a rough measure of the age of a galaxy or, more
precisely, its specific SFR (SSFR), since the
conversion from SSFR to age is dependent on the star formation
history of the source (e.g., Leitherer et al. 1999).
To zeroth order we may also expect the EW(Hα) to be
independent of extinction. In contrast, the complex escape path
of the Lyα photons relative to the continuum photons
combined with the effects of extinction
may result in substantial changes in the EW(Lyα).
These have been observed by some as decreases (Shapley et al. 2003;
Pentericci et al. 2009) but not by others (Atek et al. 2008;
Finkelstein et al. 2009c).
We note, however, that even the EW(Hα)
may have some dependence on extinction, since the stars
producing the ionizing photons may lie in different spatial
regions of the galaxy than those producing the optical continuum
photons.

Figure 10: The rest-frame EW(Lyα) vs. the rest-frame EW(Hα).
The red squares show the LAE galaxies (the error bars are ±1σ),
including those with rest-frame EW(Lyα) between 15 and 20 Å from the supplement to Table 1.
The blue diamonds show the UV-continuum sample placed at the average
EW(Lyα)=7 Å that was measured from the stacked spectrum.
The positioning of these objects in y is purely for display
purposes and is not used in the analysis.
The black line shows a linear relation corresponding to the median
EW(Hα)/EW(Lyα)=3.5 that was measured from the full LAE
sample with rest frame EW(Lyα)>15Å.

Figure 11: The (a) SDSS u′−z′ color, (b) FWHM
in CFHT MegaPrime U-band images, (c) UV power law slope,
(d) f(Hβ)/f(Hα),
(e) f(Lyα)/f(Hα), and
(f) log(f([NII]λ6584)/f(Hα)),
all plotted vs. the rest-frame EW(Hα).
The red squares show the LAE galaxies with EW(Lyα)≥20 Å.
The blue diamonds show the UV-continuum sample.
Error bars are ±1σ.
In (a) we have used the spectra to correct for the emission-line
contributions to the z′ magnitude (see Section 5.2).
In (a) and (c) the black curves show solar metallicity models
from STARBURST99 with no extinction: a constant SFR model
(thick solid) and a single instantaneous starburst model
(dashed).
The right-hand axes in (d) and (e) show
the E(B-V) values that would be derived from the Calzetti et al. (2000)
reddening law. The right-hand axis in (f) shows the O abundance
that would be derived from f([NII]λ6584)/f(Hα)
using the Pettini & Pagel (2004) conversion.

While in the absence of dust we may expect the flux of Lyα to be related
to that of Hα by the case B ratio (8.7 for a temperature of 10,000 K;
Brocklehurst 1971),
the relation between the EW(Hα) and the EW(Lyα) is more complex
because of the relative evolution of the continuum fluxes. The EW(Lyα)
declines rapidly after a single burst of star formation, as does
the EW(Hα). However, for a constant SFR model, the EW(Lyα)
instead becomes roughly constant, since the UV-continuum flux and the Lyα
line strength are governed by ongoing massive star formation
(Charlot & Fall 1993; Schaerer & Verhamme 2008).
In contrast, the EW(Hα) continues to decline, even for a constant SFR model,
as the older stars build up in the galaxy. For a galaxy with a large old population
in place, the onset of a new burst of star formation will give a high EW(Lyα)
while leaving the EW(Hα) relatively weak. It is this ongoing build-up of the
older stars that makes the EW(Hα) a function of the age of the source.

In Figure 5.1 we show the rest-frame EW(Lyα) versus the
rest-frame EW(Hα) for the LAE galaxies (red squares),
including those with rest-frame EW(Lyα) between 15 and 20 Å from the supplement to Table 1.
We also show the spread of the EW(Hα) for the
UV-continuum sample (blue diamonds).
While there is a large scatter, there is a
broad general trend for the LAEs to have higher EW(Hα). Specifically, we
find that in the EW(Lyα)≥20 Å and EW(Lyα)=15−20 Å samples, the median EW(Hα) is 113 (66−233) Å and 71 (47−77) Å,
respectively, while in the UV-continuum sample, it is 38 (31−44) Å,
where the quantities in parentheses are the 68% confidence range.
A Mann-Whitney rank-sum test gives less than a 3×10−6 probability
that the distribution of EW(Hα) is the same in the LAE
sample as in the UV-continuum sample. For the 12 highest EW((Lyα)
objects with rest-frame EW(Lyα)≥40 Å the rank sum
test give less than a 8×10−3 probability that the distribution of EW(Hα)
is the same as that for those with with weaker EWs between 15 and 40 Å.
This suggests that the LAEs are preferentially
drawn from the younger galaxies in the sample.

These results differ from those of Östlin et al. (2009), who suggested,
based on their observations of local blue compact galaxies, that there
might be an anti-correlation between EW(Lyα) and EW(Hα).
As Östlin et al. noted, their small sample is highly selected and their
result is not statistical. The Östlin et al. sample does emphasize, however,
that there are galaxies such as SBS 0335-052 that
have a very high EW(Hα)=1434 Å but
do not have strong Lyα emission.
We do not detect any sources this extreme in
our UV-continuum sample, where only four sources even have
EW(Hα)>100 Å, so they are clearly rare.

We can take this further by comparing other properties
of the galaxies with the EW(Hα). In Figure 11
we plot the following versus the rest-frame EW(Hα):
(a) the SDSS u′−z′ colors, (b) the FWHM of the
galaxies in the CFHT U-band images, (c) the UV
spectral index, (d) the flux ratio f(Hβ)/f(Hα),
(e) the flux ratio f(Lyα)/f(Hα), and
(f) the logarithm of the flux ratio f([NII]λ6584)/f(Hα).
Panel (b) is restricted
to the subsample covered by the U-band CFHT images, while panels
(a), (d), and (e) correspond to the subsample covered by the SDSS
observations. Panels (c) and (f) contain the full sample.
In panel (a) we have used the spectra to correct for the contributions
of emission lines to the z′ magnitude (see Section 5.2), since
these can make significant contributions that are a function of the
EW(Hα). We now consider each of these relations in detail.

5.2 Colors and Ages

The u′−z′ colors of Figure 11(a) follow a single
well-defined track (see also Overzier et al. 2008, who show a
similar figure for their low-redshift analogs to the LBGs).
We compare the data with a constant star formation model (solid curve)
computed with the STARBURST99 code (Leitherer et al. 1999).
Here the EW(Hα) decreases as the age increases so time
increases to the left in the plot.
The predicted evolution of the u′−z′ color versus the EW(Hα)
up to an age of 109 yr matches to the emission-line corrected
colors. By contrast, a single starburst (dashed curve) does not match,
since the EW(Hα) drops rapidly with time, while the u′−z′
colors remain blue.

We can use the full color information to infer ages and extinctions
by spectral energy distribution (SED) fitting.
For sources with strong optical emission lines
we emphasize that it is critical to remove the substantial
line contributions. The correction is a function of the EW(Hα)
and hence introduces a systematic bias into any comparison of LAEs with
UV-continuum selected galaxies. For strong-line sources SED fitting to
the uncorrected data substantially overestimates the ages. This then
causes the extinctions to be underestimated and the masses to be
overestimated.

We illustrate the problem in Figure 5.2, where we show
the broadband fluxes before (black open squares) and after (red diamonds)
line correction for three LAE galaxies spanning a wide
range in the EW(Hα). The strong [OIII]λ5007 and Hα line
contributions raise the SED at wavelengths above 4000 Å. In the fitting
procedure this is spuriously interpreted as a strong 4000 Å break,
which results in a substantial age overestimate. The strongest
emitters (Figures 5.2(a) and 5.2(b)) are, in fact,
very young but are misinterpreted as having ages of 0.45 and 0.9 Gyr
when fits to the uncorrected data are used.

We have fitted ages and extinctions for all of
the galaxies covered by the SDSS in both the LAE
and UV-continuum samples. We used the Bruzual & Charlot
(2003) models with solar metallicities and a range
of exponentially declining SFR models with e-folding
times from 1 Gyr to 20 Gyr. We combined these with a Calzetti
et al. (2000) extinction law. For each star formation
history we determined the age
and extinction that produced a χ2 minimized
fit using the measured errors in the broadband fluxes,
together with a 10% error to allow for systematic
uncertainties.

It is well known that the EW(Hα) can be used as an age
indicator (e.g., Leitherer et al. 1999; Bruzual & Charlot 2003).
We plot the galaxy ages determined from the SED fits versus
the EW(Hα) in Figure 5.2.
In Figure 5.2(a) we show the spurious ages
that would be derived from a simple fit to the
data without the emission-line corrections.
In Figure 5.2(b) we show the ages derived from
the corrected broadband fluxes.
The inferred ages are somewhat degenerate with
the assumed exponential star formation history, so
we also show the range of ages
corresponding to the range of SFR models to give a measure
of the modeling uncertainties.
We can see from Figure 5.2(a) that
nearly all of the sources with strong EW(Hα)
are assigned ages of about a Gyr. Thus, without
the line correction, we would consider there to
be no difference in age for a wide range of
EW(Hα). However, we can see from Figure 5.2(b)
that when the line correction is made, we obtain a smooth
evolution in age with the strongest EW(Hα) sources
having the lowest ages of about 10 Myr. The high EW(Hα)
sources are quite well described by a constant
SFR model (black solid curve in Figure 5.2(b)),
but at lower EW(Hα), the sources tend to fall below this
curve, presumably as the SFR begins to decline. The data are quite
well represented by a single power law fit

logT=5.48±0.23+(−1.53±0.12)×log(EW(Hα)),

(1)

where the age T is in Myr and the EW(Hα) is in Å.
This is shown by the green dashed line in Figure 5.2(b).

Figure 12:
SEDs before (black open squares) and after (red solid diamonds) removal
of the emission-line contributions. Each panel shows a LAE
galaxy with the rest-frame EW(Hα) given in the upper-left corner.
For each spectrum we removed the Lyα, [OII]λ3727, Hγ,
Hβ, [OIII]λ4959 and λ5007, Hα, [NII]λ6548 and
λ6584, and [SII]λ6716 and λ6738 emission lines using the
Gaussian fits. We then integrated the spectrum (with and without the lines) through
the filter response to determine the fraction of the light that is produced
by the continuum and used this to correct the magnitudes. The principal
contributing lines are marked at the top of the figure. The error bars are
±1σ, including a 10% flux error to allow for possible
systematic errors in the determination of the total magnitudes. The line
effects are the largest for the higher EW(Hα) galaxies. The
corrections act to smooth the SEDs. For each galaxy we show the constant
SFR model fits from Bruzual & Charlot (2003). The black (red) curve
shows the fit to the raw (corrected) data.
Both are reduced by a factor of three so the data points can be clearly
seen. The fitted ages are shown in the upper left corner for the raw fit
(black) and for the line-corrected fit (red).

Figure 13: The ages determined from spectral synthesis fitting to the
Bruzual & Charlot (2003) models vs. the EW(Hα). In (a) we show
the ages derived when no correction is applied for the emission lines.
In (b) we show the ages derived when the broadband colors are properly
corrected. For each object
(red squares—LAE galaxies with EW(Lyα)≥20 Å;
blue diamonds—UV-continuum sample)
we show the range of ages corresponding
to the models with exponentially declining SFRs with e-folding times
from 1 Gyr to 20 Gyr (vertical bars). The data point
corresponds to the 20 Gyr case.
In (b) the black solid curve shows the age vs. the
EW(Hα) relation for a constant SFR model computed with
STARBURST99, and the green dashed line shows the optimal power law
fit given in the text.

Figure 14: The surface density of sources vs. the EW(Hα). The red
shaded histogram shows the LAE galaxies with rest-frame
EW(Lyα)≥20 Å. The blue histogram shows the UV-continuum sample,
which is much more heavily weighted to low values of the EW(Hα).
The black solid curve shows the shape computed from a
model where we assume a constant rate of formation of galaxies and
that the time spent in each equivalent width bin is determined
from Equation 1. The normalization is matched to the observed surface
density.

In Figure 5.2 we plot the surface
densities of galaxies in our LAE and UV-continuum
samples versus the EW(Hα). It is worth
emphasizing that in both the LAE and UV-continuum samples the galaxies
are chosen solely on the basis of their NUV
magnitudes; that is, they have NUV<22.1, lie in the redshift
interval z=0.195−0.44, and have the same NUV magnitude distribution.
The only difference between the two samples is the presence or absence
of strong Lyα emission. Thus, while there
might be a relation between the NUV magnitude and the EW(Hα),
this will operate in the same way in both samples and will not
affect any comparison.

For the LAE sample (red shaded histogram) we plot the sum of the inverse
areas over which a galaxy with a given NUV magnitude could be observed
in the GALEX fields (Cowie et al. 2010).
(We also corrected the LAE density for the 10% of sources that were
not spectroscopically observed.) For the UV-continuum
sample (blue histogram) we also weight the areas with the fraction of
sources that we observed as a function of NUV magnitude.
The black solid curve shows the shape computed from a
model where we assume a constant rate of formation of galaxies and
that the time spent in each equivalent width bin is determined
from Equation 1. The normalization is proportional to the
birthrate of galaxies, and we have matched it to the observed surface
density. We can see that this model matches well to the shape
at high EW(Hα), so the surface density of galaxies versus the
EW(Hα) is also consistent with a constant production rate of
new galaxies that then evolve with a constant SFR.

We can also see from the figure that Lyα emission is rare
in sources with EW(Hα)<100 Å (about 0.7% of galaxies) but
common in higher EW(Hα) sources where the blue and red histograms
become comparable. 31±13% of the EW(Hα)>100 Å galaxies
and 57±30% of the EW(Hα)>250 Å galaxies have Lyα
emission with the EW(Lyα)≥20 Å. The uncertainties reflect the
small number of UV-continuum sources with high EW(Hα). Finally,
we can see that a large fraction of the LAEs (75±12%)
are drawn from the high EW(Hα)>100 Å population. This again
shows that the LAEs are drawn primarily from the youngest galaxies.
The remaining 25% of LAEs with EW(Hα)<100 Å may have geometries, kinematics, or orientations that are unusually
conducive to Lyα escape. However, some may be objects where the
Lyα emission is produced by AGN activity, but the AGN signatures
in the optical are swamped by the galaxy contributions.

In the following sections we will use the EW(Hα)=100 Å as a
rough cut above which Lyα emission is more common.

5.3 Sizes

Figure 11(b) shows that there
is also a size evolution as a function of EW(Hα).
The higher EW(Hα) sources are generally unresolved
at the ∼1′′ resolution of the CFHT MegaPrime U-band images,
while the lower EW(Hα) sources are generally extended.
However, unlike in the other correlations, there is a significant
difference between the LAEs and the UV-continuum sources at the same
EW(Hα). From Figure 11(b) we can see that at the
same EW(Hα), sources with Lyα emission are significantly
smaller than those without; that is, where Lyα is strong is in the
most compact sources at a given EW(Hα). A rank-sum test shows that
there is a <3.6×10−5 probability that the LAEs
at EW(Hα)>80 Å are physically as large as the UV-continuum
sources with the same EW(Hα) selection. (We have chosen the
EW(Hα) to provide a significant number of UV-continuum sources
since there are only 4 UV-continuum sources with EW(Hα)>100 Å as can be seen in Figure 11(b)). The
median FWHM for these LAEs is 1.1′′, comparable to
the resolution in the CFHT MegaPrime images, while that for the
UV-continuum sources with the same EW(Hα) selection is 1.6′′.

Figure 15: Images from the GEMS survey of the CDF-S.
The lower 12 sources
are UV-continuum sources in the region, and
the top 4 sources are LAEs. Each thumbnail is
6′′ on a side. The blue and green colors correspond to the
F606W GEMS image, and the red color corresponds to the F850W
GEMS image.

Figure 16: Contour plots for the 4 LAEs in the GEMS survey of the
CDF-S. Each contour in the combined F606W and F850W image
rises by a factor of two. The rest-frame EW(Lyα) and the
rest-frame EW(Hα) are given in the lower-left corner.

The ground-based U-band imaging is inadequate to resolve the
sizes for most of the LAEs. Thus, although the optical
morphologies may differ from those in the UV, here
we use HST-based images to obtain rest-frame optical sizes.
In Figure 5.3
we show the combined F606W and F850W HST ACS images for the
sources that lie in the GEMS survey field. In Figure 5.3
we show the contours from the combined F606W and F850W images
for the four LAEs. The FWHM in these higher resolution images
is 0.5 kpc for the three sources with EW(Lyα)≥20 Å, though
the one source with EW(Lyα)<20 Å is more extended
with a FWHM of about 2 kpc. These sizes are very similar
to the continuum sizes of LAEs in the rest-frame UV at
3<z<6.5 (Bond et al. 2009, 2010; Cowie et al. 2011)
though again we wmphasize that it would be desirable to make the comparison
at the same rest-frame wavelength.

5.4 Extinction

We can measure the dust extinction from the galaxies in several
ways: from the UV spectral slopes, from the SED fits, and from
the Balmer ratios. We may intercompare these measurements
and also use each of them to test the effects of
extinction on the Lyα line.

In Figure 11(c) the UV spectral index,
which is often used as a measure of the stellar extinction,
shows a general trend towards less negative values as
we move to lower EW(Hα), which would be expected
if the extinction is rising as the galaxies age. However,
there is a good deal of scatter, and at least part
of the change is caused by evolution in the intrinsic spectrum.
The evolution of the reddening may be more clearly
seen in Figure 11(d):
the Balmer ratio shows a systematic decrease
as the EW(Hα) decreases, indicating an increase in
the nebular extinction. The derived E(B-V) is similar whether
we use the Cardelli et al. (1989) extinction law or
the Calzetti et al. (2000) extinction law. With the Calzetti et al. reddening we obtain a median E(B-V) of 0.23 (0.15, 0.32)
from the Balmer ratio for sources with EW(Hα)>100 Å and 0.33 (0.30, 0.38) for sources with EW(Hα)=20−100 Å,
where the quantities in parentheses are the ±1σ range.
Interestingly, in both Figures 11(c) [UV spectral index versus
EW(Hα)] and 11(d) [Balmer ratio versus EW(Hα)]
there appears to be little differentiation between LAEs and UV-continuum
galaxies at the same EW(Hα), suggesting that extinction
is not the primary reason for Lyα being suppressed at a
given EW(Hα). This is a well-known result
which has been found by many authors (e.g., Giavalisco et al. 1996; Mas-Hesse et al. 2003;
Atek et al. 2009b; Finkelstein et al. 2009a).

Figure 17: The evolution of (a) f(Lyα)/f(Hα) and
(b) the UV spectral index vs. the Balmer ratio
f(Hα)/f(Hβ).
All sources with rest-frame EW(Hα)>20 Å are shown (red squares — LAEs; blue diamonds — UV-continuum
sources). The EW(Hα) limit is chosen to include
nearly all of the LAEs (see Figure 5.2) and to provide
a sample of UV-continuum objects with strong Hα.
The two figures contain slightly different sets of objects
since the UV spectral indices are only measured for a limited
redshift range.
The upper limits on the Lyα fluxes for the UV-continuum
galaxies are computed for an observed-frame EW(Lyα)=10 Å.
The expected ratios are shown for several reddening laws:
Calzetti et al. (2000; solid), Cardelli et al. (1989; dashed),
and Fitzpatrick et al. (1999; dotted).
The crosses on the curves show increments of 0.1 in E(B-V) from the
adopted intrinsic values, which are shown by the solid vertical
and horizontal lines. In (a) the lower (upper) green line shows
the expected relationship for the Calzetti et al. reddening law
if the E(B-V) is twice (half) as large for Lyα as for the
Balmer lines.

As has been pointed out by Atek et al. (2009a) and
Scarlata et al. (2009) based on the GALEX LAE samples
(see also Hayes et al. (2010) for a discussion of a higher
redshift sample), there is a relation between
f(Lyα)/f(Hα) and f(Hα)/f(Hβ)
for LAE galaxies but with a wide dispersion.
In Figure 5.4(a) we show such a plot for our LAE galaxies
(red squares) and UV-continuum galaxies (blue diamonds) with
EW(Hα)>20 Å, where the Hα limit is chosen
to include nearly all the LAEs (see Figure 5.2).
The upper limits on the Lyα fluxes for the UV-continuum
galaxies are computed for an observed-frame EW(Lyα)=10 Å.
We emphasize again that, unlike Atek et al. (2009a), our
Hα measurements are not matched to the GALEX aperture,
so our f(Lyα)/f(Hα) ratio may have additional
systematic scatter.
In addition, the relative flux calibration
and uncertainties in the underlying absorption correction
for the Hβ line may result in scatter in f(Hα)/f(Hβ)
and occasional unphysical values (negative extinction) in this ratio.

The inclusion of the UV-continuum sample shows that the
LAEs correspond to the maximum values of f(Lyα)/f(Hα)
at a given f(Hα)/f(Hβ), as would
be expected from the initial selection in the Lyα
line. At any given f(Hα)/f(Hβ) there is
a range in f(Lyα)/f(Hα) stretching up to
roughly the value set by simple extinction of the
Lyα line if it followed the same escape path as
the Hα line. The LAEs primarily lie along this
relation. The line followed by the LAEs is consistent either with the
Calzetti et al. (2000) reddening law (black solid line)
or with uniform screen models, such as those of
Cardelli et al. (1989; black dashed) or
Fitzpatrick et al. (1999; black dotted), all of which give
nearly identical predictions. There is considerable scatter
about the relation, but we suspect that this is a consequence
of systematic errors in the flux calibration discussed above, rather
than an indication of the need for more complex dust models,
as suggested by Scarlata et al. (2009). We can also see that
the data are not consistent with the Lyα photons in
the LAE galaxies having
a widely different E(B-V) path from the Hα photons.
Following Scarlata et al.,
we have plotted the relation when the Lyα
path is twice that of the Balmer photons (lower green line)
and also when it is half that of the Balmer photons (upper
green line). It is clear that these are a much poorer fit to
the LAE observations.

In contrast, if we plot the UV spectral index for the same
sample versus f(Hβ)/f(Hα) (Figure 5.4(b)),
then we see that
the Calzetti et al. (2000) reddening law gives a substantially
different result than the Cardelli et al. (1989) or
Fitzpatrick et al. (1999) uniform screen models, reflecting their
different shapes at UV wavelengths. Most of the data appear to
follow the models of Cardelli et al. or Fitzpatrick et al. rather than the reddening law of
Calzetti et al., though there are a handful of outlying points.
This suggests that for these galaxies the stellar extinction is better
reperesented as a uniform screen rather than a patchy distribution.
This is the case for both the LAEs and the UV-continuum sources.

5.5 Metallicity

Perhaps the most interesting plot in Figure 11 is
that of N2=log(f([NII]λ6584)/f(Hα)) versus
EW(Hα) (panel f). N2 is a widely used metallicity indicator
with low values corresponding to low metallicities. The fact that
both the LAE and UV-continuum galaxies follow a surprisingly tight
track in this diagram with N2 rising
as EW(Hα) decreases appears to show that we are seeing a
smooth metal build-up with age in both classes of galaxies.

In order to translate N2 to
a metallicity in the galaxies we use the Pettini & Pagel (2004) local
relationship, 12+log(O/H)=8.90+0.57N2, which was determined
by comparing to direct measures of the O abundance over the range
N2=−2.5 to −0.5. Extrapolating the Pettini & Pagel relation to these
galaxies requires assuming that there is no change in the ionization
parameter, which may well be incorrect. However, at z=0.195−0.44
Cowie & Barger (2008) found a narrow range of ionization parameters
(q∼2×107) and a similar relation to that of Pettini & Pagel.
Cowie & Barger (2008) also showed that other line diagnostics
gave similar metal-luminosity relations to that derived from N2 in the
redshift interval z=0.195−0.44.

Figure 18: 12+log(O/H) derived from the direct method vs. N2.
The GALEX sources in the redshift interval z=0.195−0.44
are shown with red squares with 1σ error bars.
A sample of ultra-strong emission-line galaxies in the
same redshift interval with direct O measurements from Hu et al. (2010)
are shown with green diamonds. Both are broadly consistent
with the Pettini & Pagel (2004) relation (black line), though
with substantial scatter.

We can also test the Pettini & Pagel (2004) relation with the present
data. Ten of the z=0.195−0.44 sources have [OIII]λ4363
detected at above the 3σ level. These are marked with
green diamonds in Figure 3(a). None of the UV-continuum sample
at z=0.195−0.44 has detected [OIII]λ4363 at the 3σ level.
All of the [OIII]λ4363 galaxies have low values of N2, except for
GALEX 1240+6233, which is classified as an AGN based on the BPT
diagram. Conversely, most of the low-N2 galaxies are detected in
[OIII]λ4363. For the sources that are not AGNs and where we
have detected [OIII]λ4363 we used the ‘direct’ or Te method
to determine the metallicity
(e.g., Seaton 1975; Pagel et al. 1992; Pilyugin & Thuan 2005;
Izotov et al. 2006).
To derive Te[O III] and the oxygen abundances, we used the
Izotov et al. (2006) formulae, which were
developed with the latest atomic data and photoionization models.
We compare the derived O abundances with N2 in Figure 5.5
(red squares), where we also show similar measurements from a sample of
ultra-strong emission-line galaxies taken from Hu et al. (2009;
green diamonds). The O abundance versus N2 is broadly consistent
with the Pettini & Pagel relation (black line), though there is a
good deal of scatter, probably reflecting the variation in the
ionization parameter.

We can see immediately from Figure 11(f) that the
LAEs in the redshift interval z=0.195−0.44 have lower values
of N2 than the UV-continuum sources
(see also Figure 3). This effect was previously
noted by Cowie et al. (2010), though the present larger sample
substantially increases the statistical significance of the result.
This result may be more clearly seen in Figure 5.5,
where we show the N2 distributions for the LAE (red shaded histogram)
and UV-continuum (blue histogram) samples. While the N2 distributions
overlap, the distribution for the LAEs clearly extends
to lower values, and the median N2 is lower. A rank-sum test
gives only a 10−6 probability that the two GALEX
samples are similar. Using the Pettini & Pagel (2004)
conversion gives a median 12+log(O/H)=8.24 (8.17, 8.35) for the
LAEs, where the quantities in parentheses are the ±1σ range.
This is about 0.4 dex lower than that of the UV-continuum sample, though
the median value for the UV-continuum sample is more poorly determined,
since at the higher values of N2 the conversion breaks down
as the N2 value saturates.

We can now see from Figure 11(f)
that the difference in N2 distributions is primarily
a reflection of the evolutionary state of the galaxy. That is, the
metallicity is building up with age, and the LAEs are preferentially
drawn from the younger galaxies. UV-continuum galaxies have
similar metallicities to LAEs with the same EW(Hα). However,
the LAE sample is much more weighted to high EW(Hα) galaxies,
which are younger and more metal poor.

Figure 19: The number distributions of the LAE galaxies
with EW(Lyα)≥20 Å (red shaded
histogram) and the UV-continuum
sample (blue histogram) as a function of N2.
The UV-continuum sample is much more heavily weighted
to high values of N2. The top axis shows the O abundance that
would be derived from the N2 ratio using the Pettini
& Pagel (2004) conversion.

In order to convert the observed quantities into physical
properties of the galaxies, we will assume a constant SFR model
with a Salpeter (1955) initial mass function (IMF). As we have
seen, this appears to provide a self-consistent fit to most of
the observed properties. For simplicity, we assume that the IMF
extends smoothly from 0.1 M\sun to 100 M\sun, though
this can be easily converted into other favored IMFs with flatter
slopes below M\sun=1 by a simple renormalization of the SFR.

In Figure 6 we show the Hα luminosities
of the LAE and UV-continuum samples corrected for extinction
using the Balmer ratio and a Fitzpatrick et al. (1999) extinction
law. The median corrections for the LAE and UV-continuum
samples are 1.7 and 2.1, respectively.

For our assumed constant SFR model the conversion from Hα
luminosity to SFR is roughly invariant over the range of EW(Hα)
and is given by the Kennicutt (1998) relation:

logSFR=−41.1+logL(Hα).

(2)

On the right-hand axis of Figure 6 we give the SFR in
M\sun yr−1 corresponding to a given Hα luminosity.
Individual galaxies have SFRs that range from about 1 M\sun yr−1
up to about 30 M\sun yr−1. The lower bound is roughly set by
the NUV magnitude limit for the selection, but the upper bound on the SFR
is a measure of the maximum SFRs seen in UV-continuum selected galaxies at
these redshifts. On the top axis we give the logarithm of
the age in years corresponding to a given EW(Hα). The LAEs
are primarily drawn from sources younger than 1−2 Gyr.

We can see from Figure 6 that there is
no evolution in the distribution of the SFRs as a function of
the EW(Hα). This is consistent with our constant SFR
assumption. There is also no difference in the distribution
of the SFRs for the LAE and UV-continuum galaxies. If the
highest SFR galaxies maintain their rates for a subtantial
fraction of the local age of the universe, then the final
stages will have stellar masses in excess of 1011 M\sun.

Figure 20: Extinction-corrected Hα luminosity vs. EW(Hα).
The red squares show LAEs with EW(Lyα)≥20 Å, and the
blue diamonds show UV-continuum galaxies. The conversion
to a SFR assuming a Salpeter (1955) IMF and a constant SFR is shown on the
right-hand axis, and the conversion to galaxy age is shown
on the top axis.

We can test our Hα luminosity-determined SFRs against those
inferred from 20 cm observations (where available; see
Section 2.1) and those determined from
the UV continuum. In Figure 6 we plot
rest-frame 20 cm power calculated assuming a radio spectrum of
fν∼ν−0.8 versus extinction-corrected Hα
luminosity. We see a nearly linear relation, though the
number of measured sources is fairly small. We compare the data with
two conversions of the radio power to SFR: Condon (1992; solid
line) and Bell (2003; dashed line). The two relations appear to
bracket the data points.
(Prior to the Hα extinction correction the points are more
scattered and fall to the left of the expected relations.)

There is one significantly deviant point, GALEX0959+0151 in COSMOS 00,
where the Hα luminosity is very high compared to the radio power.
This source has the highest EW(Hα) in the figure (811 Å) and
an age of 9 Myr from SED fitting (Figure 5.2(b)).
This suggests that the radio power is low because the supernovae that
generate the radio emission have yet to occur. Unfortunately, the
other sources in the sample with very high rest-frame EW(Hα)
are not in the fields with radio data, so we cannot test this further
at present.

Figure 21: The 20 cm power in erg Hz−1 vs. the extinction-corrected
Hα luminosity. The red squares show LAEs, and the blue diamonds
show UV-continuum galaxies. In this plot, since we are only
testing the star formation rate determinations, we have
included sources from the supplement to Table 1
with EW(Lyα) between 15 and 20 Å.
For sources not detected in the
20 cm images we have assumed upper limits of 100 μJy (SIRTFFL 01)
and 40 μJy (COSMOS 00).
These are shown with downward pointing arrows.
The one source in the HDF 00 field is detected at 20 cm.
The solid line shows the Condon (1992) SFR vs. radio power relation,
and the dashed line shows the Bell (2003)
high-luminosity SFR vs. radio power relation.

Figure 22: The ratio of the SFR computed from the extinction-corrected
UV luminosity to that from the extinction-corrected Hα luminosity.
The red squares show LAEs with EW(Lyα)≥20Å, and the blue
diamonds show UV-continuum galaxies. In computing the UV
SFRs we have applied a small multiplicative offset of 1.17 to scale
the NUV GALEX spectra to match the NUV GALEX magnitudes.
The green curve shows the evolution of the 1600 Å Fλλ
for a constant SFR model with a Salpeter (1955) IMF.

SFRs for high-redshift galaxies are generally calculated from the
observed UV luminosity together with the stellar extinction correction
determined from the UV continuum slope (Meurer et al. 1999).
It is therefore interesting to commpare the SFRs derived using
this method with those computed
from the Hα luminosities in the present low redshift
sample and to estimate the limits of its validity.
In Figure 6 we show the ratio of the SFR derived from
the stellar extinction corrected UV luminosity with that
computed from the nebular extinction corrected Hα luminosity
as a function of EW(Hα). We have computed the UV SFR from the
L1600 luminosity, extinction corrected using the value derived from
the UV spectral index, and converted using the value at t=108 yr
computed from the STARBURST99 model,

logSFR=−39.87+logL1600.

(3)

With this calibration the UV SFRs are closely matched
to the Hα SFRs with an average ratio of 0.9
for sources with EW(Hα)=60−200 Å.
This small absolute difference may be a simple consequence of
the relative optical and UV calibrations. As can be seen in
Figure 6, the UV calibration underestimates the
SFR at higher EW(Hα) and overestimates it at lower
EW(Hα).

The reasons for the deviations reflect the limit of validity of the
underlying assumptions. At low EW(Hα) the derivation of the
UV continuum extinction from the UV spectral slope breaks down.
The Meurer et al. (1999) correction is only applicable to starbursting
galaxies, since it depends on the intrinsic spectral index being
approximately fixed,
as is the case for these sources (see Figure 11(c)).
At low EW(Hα) (less than about 60 Å) the intrinsic UV
spectral indices are shallower, and the inappropriate use of the
relationship overestimates the SFR, as can be seen in
Figure 6.
A second problem with the UV continuum method is that the UV continuum flux
evolves with time, even in the constant SFR models. In
Figure 6 we show (green curve) the evolution of
Lλλ evaluated at 1600 Å (L1600) as a function
of EW(Hα). (The normalization is arbitrary and depends
on the SFR.) The UV luminosity increases as a function
of age up to ∼108 yr or an EW(Hα)∼200 Å.
Thus, the UV continuum calculation combined with the UV spectral
index extinction correction
should only be used over the EW(Hα)=60−200 Å range, or,
equivalently, for rest-frame line-corrected galaxy colors
of u′−z′=0.5−1.4 (Figure 11(a)).

Figure 23: Metallicity vs. EW(Hα). In addition to the present
data points (i.e., Figure 11(f)), we show the blue
compact galaxies from Östlin et al. (2009) (open squares:
red with Lyα emission, blue without),
a continuum-selected sample with NUV=22−23.25 from the GOODS-N
observations of Barger et al. (2008) with
z=0.195−0.44 (black triangles), and the ultra-strong Hα
emission-line selected sample from Hu et al. (2010)
(green circles). The error bars are ±1σ,
with undetected sources shown with downward pointing
arrows at the 1σ level.
The black curve shows the model
fit discussed in the text. The values of 12+log(O/H) shown on the
right vertical axis were calculated using the Pettini & Pagel (2004)
relation.

In Figure 6 we plot the metallicities of the LAE
(red solid squares) and UV-continuum (blue diamonds) samples versus
the EW(Hα).
If the gas reservoir were fixed, the constant SFR model would predict
that the metallicities should rise linearly with time in the early
stages. This is clearly much steeper than is seen in the present samples,
which can be approximated by a model where the metallicity rises as t0.3
(black curve: this is only shown over the range of validity of the
N2 diagnostic which is only appropriate for low metallicity
and saturates at near solar metallicity). Thus,
the galaxies must have ongoing accretion of gas.

In Figure 6 we also compare the present samples with a
fainter continuum sample from the GOODS-N observations of Barger
et al. (2008) (black triangles) and with a fainter sample selected
to have very high EW(Hα) from Hu et al. (2010) (green circles),
as well as with the local blue compact galaxy sample of
Östlin et al. (2009) (open squares).
As might be expected, the
fainter samples have systematically lower metallicities than the present
data. However, the effect is small compared to the EW(Hα) dependence,
suggesting that the ages of the galaxies primarily determine the metallicities.

Figure 24: (a) The distribution of the SFRs in the LAE (red shaded histogram)
and UV-continuum (blue histogram) samples. (b) The distributions of the
stellar masses.

Using the SFRs derived from the extinction-corrected Hα
luminosities and the masses derived from the constant SFR fits to
the line-corrected SEDs, we show our final result in Figure 6.
In Figure 6(a) we show the distribution of the SFRs.
The two samples are nearly identical with a median of
∼6 M\sun yr−1. In contrast, as is shown in
Figure 6(b), the LAEs, reflecting their younger ages,
are systematically lower in mass than the UV-continuum sources.
The median mass for the LAEs is 109 M\sun, while that
for the UV-continuum sample is 8×109 M\sun.

We analyzed a substantial sample of UV continuum-selected galaxies
with and without Lyα emission lines to try to understand how
LAEs are drawn from the general population and how they
evolve with redshift. To do this, we obtained extensive optical
spectroscopy of GALEX grism selected samples (both LAE samples
and comparison samples with the same UV magnitude distributions but no
detected Lyα) in two redshift intervals: z=0.195−0.44 and
z=0.65−1.25. We used the optical spectroscopy to eliminate AGNs and
to obtain the optical emission-line properties of the samples.

We confirmed that the GALEX selected z∼0.3 LAEs
are considerably fainter and much rarer than the
high-redshift LAEs (Deharveng et al. 2008; Cowie et al. 2010).
Cowie et al. (2010) showed that the
L⋆ in a Schechter (1976) function fit to the z=0.195−0.44 LAE
luminosity function is almost an order of magnitude
fainter than the L⋆ in a fit at z∼3 and that only about
5% of z∼0.3 UV-continuum selected galaxies have
rest-frame EW(Lyα)≥20 Å.
Here we showed that there are also LAEs that can be
found in the GALEX
spectra at z∼1 and that these are similar in
luminosity to the most luminous high-redshift galaxies.
Thus, we conclude that most of the observed evolution occurs over
the redshift interval z=0−1.
This appears to be a simple downsizing effect, with the presence
of the higher Lyα luminosity sources at z=1
corresponding to higher SFR sources initiating their
star formation at this redshift.

We showed that at z=0.195−0.44, SED fits to strong emission-line galaxies
significantly overestimate the ages and masses and underestimate the
extinctions if we do not correct for the emission-lines. We found that the SED fits to the uncorrected
broadband fluxes give ages of about a Gyr. However, the inferred ages drop
substantially when we correct the galaxy broadband magnitudes using
the observed spectra to remove the emission-line contributions.
All the galaxies with ages much less than a Gyr have strong emission lines
and must be corrected for the line contributions.
Spectral synthesis fitting shows a smooth evolution of the rest-frame
EW(Hα) with age when the emission-line
contributions are removed from the broadband fluxes.

We found that at z=0.195−0.44, all sources, regardless of the
strength of the Lyα emission line, follow a single,
well-defined sequence as a function of the rest-frame EW(Hα).
Higher EW(Hα) sources all have lower metallicities,
bluer colors, smaller sizes, and less extinction.
The number distribution of galaxies versus the EW(Hα) is consistent
with a constant formation rate of new galaxies.
The bulk (75±12%) of the LAEs lie at high EW(Hα)>100 Å, and
31±11% of all UV-continuum selected galaxies with EW(Hα)>100 Å are LAEs.
We conclude that the low-redshift LAEs are primarily drawn from a population of
young galaxies that have recently initiated star formation.

It appears that LAEs represent an early stage in a starburst
when the star-forming gas is still relatively pristine and
the primary star-forming region is small. It also appears that
there is a time sequence, with the Lyα emission line dying
away and the metallicity of the gas rising as the galaxy evolves.

We are indebted to the staff of the Keck observatory for their
excellent assistance with the observations. We would like to thank
Toni Songaila and the anonymous referee for their critical reading of the paper and
useful suggestions for improving it.
We gratefully acknowledge support from NSF
grants AST-0709356 (L. L. C.), AST-0708793 (A. J. B.),
and AST-0687850 (E. M. H.), from the University of
Wisconsin Research Committee with funds granted by the
Wisconsin Alumni Research Foundation and from the David and
Lucile Packard Foundation (A. J. B.), and from
a NASA grant through an award issued by JPL 1289080 (E. M. H.).

Appendix A Appendix

Since the analysis of Cowie et al. (2010), much deeper GALEX
grism spectroscopic observations were released for the CDFS 00 field.
We have obtained the one- and two-dimensional GALEX spectra from
MAST and analyzed them to search for Lyα emission using the same
procedures as in Cowie et al. (2010). We find a sample of 100 Lyα
selected sources within a 32\farcm5 radius field.
We give the properties of this sample in Table A1. For each source
we give the number, the GALEX name, the J2000 right ascension and
declination in decimal degrees, the NUV and FUV magnitudes, the redshift
inferred from the Lyα emission line in the GALEX UV
spectrum, the line width in km s−1 together with the 1σ error,
whether the galaxy is classified as an AGN based on the presence of
high-excitation lines in the UV spectrum, and, finally, the optical redshift,
if available. The optical redshifts are primarily taken from our DEIMOS
spectroscopy, but we also include redshifts from Balestra et al. (2010)
and Vanzella et al. (2008), which we note with a colon after the
redshift. Where a source lies within the X-ray observations of the
Extended Chandra Deep Field-South (ECDF-S) and is (is not) detected in
X-rays (Lehmer et al. 2005), we give the logarithm of the rest-frame
2−8 keV luminosity (an “E”) in parentheses in the optical redshift
column (col. 10).